Cover Page
The handle http://hdl.handle.net/1887/78122 holds various files of this Leiden University dissertation.
Author: Vardanyan, V.
A C C E L E R AT I O N
Proefschrift
ter verkrijging van
de graad van Doctor aan de Universiteit Leiden,
op gezag van Rector Magnificus prof. mr. C.J.J.M. Stolker,
volgens besluit van het College voor Promoties
te verdedigen op woensdag 18 September 2019
klokke 10:00 uur
door
Valeri Vardanyan
Promotor: Prof. Dr. Ana Achúcarro Co-promotor: Dr. Alessandra Silvestri
Promotiecommissie: Prof. Dr. L. Amendola (ITP, University of Heidelberg, Germany) Prof. Dr. A.-C. Davis (DAMTP, University of Cambridge, UK) Prof. Dr. E. R. Eliel
Prof. Dr. H. Hoekstra
The cover shows a L-shaped sketch of W. de Sitter. The Dutch text around the balloon translates into "Who is really blowing up the balloon? What makes the universe expand, or swell up? This is done by Lambda. Another answer cannot be given". The original version of the sketch first appeared accompanying an interview with De Sitter in Algemeen Handelsblad, pub-lished in 1930. The large scale structure image is taken from the website of The Millennium Simulation Project. The author of this thesis thanks Tsovinar Martirosyan for realizing the idea.
Casimir PhD series, Delft-Leiden 2019-31 ISBN 978-90-8593-414-1
An electronic version of this thesis can be found athttps://openaccess.leidenuniv. nl
C O N T E N T S
1 introduction 1
1.1 The cosmological standard model in a nutshell . . . 4
1.2 Observations . . . 15
1.3 The inflationary paradigm . . . 19
1.4 Beyond the standard model . . . 22
1.5 Screening mechanisms in modified gravity . . . 31
1.6 The era of precision cosmology . . . 33
1.7 This thesis . . . 35
i cosmic acceleration in supergravity 2 dark energy, a-attractors, and large-scale struc-ture surveys 41 2.1 Introduction . . . 43
2.2 Asymmetric cosmological a-attractors . . . . 47
2.3 a-attractors and supergravity . . . . 56
2.3.1 General formulation, geometry, and special values of a . . . . 56
2.3.2 Suppressing the fifth force . . . 59
2.4 Single-field quintessential inflation models . . . 61
2.4.1 Inflationary dynamics, late-time evolution, and cosmic accel-eration . . . 61
2.4.2 Gravitational reheating versus instant preheating . . . 67
2.4.3 Spectral index: Comparison with the non-quintessence scenario 70 2.5 Examples of single-field models of quintessential inflation . . . 73
2.5.1 Linear potential . . . 73
2.5.2 Two-shoulder model with exponential potential . . . 81
2.5.3 Exponential potential . . . 84
2.6 2-field quintessential inflation models . . . 104
2.6.1 Dark energy and exponential potentials . . . 104
2.6.2 Non-interacting a-attractors . . . . 106
2.6.3 Interacting a-attractors . . . . 110
2.6.4 Quintessence with a linear potential . . . 112
2.6.5 Comparison to observations, and constraints on parameters . 114 2.7 Conclusions . . . 117
ii modifications of gravity 3 doubly coupled bigravity and the gw170817 event 123 3.1 Introduction . . . 124
3.2 Cosmology of doubly-coupled bigravity . . . 127
3.3 The speed of gravitational waves . . . 132
3.4 MCMC scans and observational constraints . . . 142
3.4.1 One-parameter models . . . 145
3.4.2 Two-parameter models . . . 151
3.4.3 Further remarks . . . 154
3.5 Conclusions . . . 156
3.6 Appendix: Tensor modes . . . 158
4 mimetic massive cosmology 171 4.1 Introduction . . . 172
4.2 Flat-space perturbations . . . 176
4.3 Cosmological solutions . . . 180
4.4 Cosmological perturbations . . . 184
4.4.1 Stability bound . . . 185
4.4.2 Cosmological tensor mass . . . 190
4.4.3 Quasistatic limit . . . 192
4.5 Conclusions . . . 193 iii screening mechanisms in modified gravity
contents vii
5.1 Introduction . . . 201
5.2 Spherical Collapse . . . 203
5.3 Self-similar spherical collapse . . . 206
5.4 Symmetron gravity . . . 212
5.5 Spherical collapse with the Symmetron . . . 216
5.5.1 Field profile . . . 217
5.5.2 Splashback . . . 221
5.6 Discussion and conclusion . . . 227
1
I N T R O D U C T I O NCosmological research is about the global, large-scale properties of the universe. It is one of the most actively developing fields of modern physics. This rapid flourishing of the field is partly motivated by the Nobel Prize winning discovery of cosmic acceleration in 1998 [1,2], and partly by the fact that cosmology can serve as a uniquely fascinating laboratory for testing various aspects of fundamental theories of physics. Indeed, it is already widely acknowledged that the cosmological observations suggest tests at regimes which are by far not accessible at the laboratory setups.
All the wealth of cosmological observations are consistently explained by a phenomenological model referred to as the cosmological standard model. This model assumes, first of all, that the universe is homogeneous and isotropic at the largest scales. Additionally, it is now well measured that the biggest share in the energy budget of the universe, about 68%, belongs to the cosmological constant, L - a constant energy density component with a negative pressure. Such a component causes the universe to expand with increasing rate, a phenomenon known as cosmic acceleration. In addition to this, about 27% of the universe is composed of a non-relativistic, pres-sureless gas called cold dark matter, which interacts gravitationally, but does not interact electromagnetically, and hence can be observed only through its gravitational effects. The conventional baryonic matter and radiation together make only about 5% of the universe’s energy budget. This matter
content, together with a hypothesized short period of very rapid expansion of spacetime in the very early universe, known as cosmic inflation, provides a beautifully simple interpretation of practically all the currently available cosmological observations in the context of General Theory of Relativity (GR). This cosmological model is often referred to as the L-Cold Dark Matter (LCDM) model.
The rough timeline of the universe is that it experienced a rapid (infla-tionary) expansion during its earliest stages. This expansion caused most of the inhomogeneity and anisotropies in spacetime to reduce, and the spatial curvature to flatten out (see later in this chapter). After the inflationary stage the universe reheats, i.e it becomes dominated by a relativistic plasma. As universe expands, the energy density of this relativistic plasma decreases and the universe enters the epoch dominated by nonrelativistic particles -baryons and dark matter. At some point the energy of collisions in cosmic plasma decreases so much that neutral atoms are formed, and the residual photons, unable to Compton-scatter on free electrons anymore, freestream through the entire universe. Later on, as the universe becomes dominated by dark matter, the small fluctuations in density start to grow, eventually leading to formation of galaxies and galaxy clusters. The matter dominated epoch then is followed by an accelerated expansion caused by yet unknown mechanism. Phenomenologically the simplest candidate for this unknown mechanism is the cosmological constant mentioned above.
introduction 3 It is a very curious fact that the standard framework of quantum field theory already leads to accelerated expansion of the universe. Indeed, quantum mechanically we expect a non-zero vacuum energy, which behaves exactly like a cosmological constant. If the theoretically estimated value of the vacuum energy density would agree with the cosmological observations, this would have been one of the most elegant predictions in theoretical physics. Unfortunately the reality is by far not as simple as that. The trouble is that the theoretical expectation for the value of this vacuum energy is at least tens of orders of magnitude larger than the value inferred from cosmological observations (see [3] for a pedagogical treatment of the topic). Besides the quantum mechanical contribution, there is also a classical contribution to the vacuum energy density originating, e.g., from the minima of scalar field potentials. The huge value of the quantum mechanical vacuum energy can in principle be cancelled against the classical contributions. This cancellation between two huge values, however, is highly unsatisfactory as we would need a very precise, finely-tuned cancellation.
The line of research of exploring the alternatives to the cosmological stan-dard model, while originating from the need of explaining the accelerated expansion, has now to some extent diverged from its origins. Indeed, now a big part of research in this direction is devoted to using cosmological observations for testing various theoretical models, without necessarily re-quiring these models to give cosmic acceleration in absence of cosmological constant.
The theme of this dissertation is largely motivated by the phenomenon of cosmic acceleration and is devoted to understanding various properties of the fundamental laws of nature by exploiting the cosmological phenomena. Before moving to the main chapters of this thesis, let us quickly review the main ideas in modern cosmology.
1.1 the cosmological standard model in a nutshell Homogeneous and isotropic universe
In order to understand the basic dynamical properties of the universe, we should note that the most relevant interaction at such large scales is the gravity. Our current picture of the latter is dominated by the fact that spacetime is a dynamical object described by the metric tensor gµn (we use
Greek indices for denoting the 4-dimensional spacetime coordinates). In this thesis we will employ the ( ,+,+,+) sign convention for the metric.
Cosmological observations suggest that on very large scales (larger than
O(100)megaparsecs) the universe is described by a spatially homogeneous
and isotropic manifold, first presented by Friedmann [4,5]. The most gen-eral metric compatible with spatial homogeneity and isotropy is known as Friedmann-Lemaître-Robertson-Walker (FLRW) metric and can be written as
ds2 = N(t)2dt2+a(t)2✓ dr 2
1 kr2 +r
1.1 the cosmological standard model in a nutshell 5 where t is the time coordinate, r is a radial coordinate on the spatial hyper-surfaces, d2W is the metric of a two-sphere and k is introduced for
account-ing for the spatial curvature of the metric. As we see, we need to introduce two functions of time, N(t)and a(t)known as the lapse function and the scale
factor of the universe. The former is related to the time-reparametrization invariance of the metric, and can be safely fixed to any functional form. This reparametrization invariance originates from the symmetries of General Theory of Relativity to be discussed below. Two important choices for N(t)
are the so-called cosmic time, corresponding to N(t) = t and the conformal
time, corresponding to N(t) = a(t). The scale factor keeps track of how
length intervals on spatial slices of spacetime shrink or expand over cosmic time t. For example, the ratio of physical distances between two galaxies at times t1 and t2 is simply given by a(t1)/a(t2). This change between
the distances is an inherent feature of an FLRW metric and should not be confused with the change caused by the peculiar motion of galaxies, which can be, for example, due to the gravitational force exerted on the considered galaxies by their neighbouring mass. An additional comment on terminology is appropriate here. The radial coordinate r in FLRW metric is typically referred to as a comoving coordinate. This reflects the fact that the distance r between two point does not change during the cosmic evolution. The physical distance between two points, rphys = a(t)r, however, of course
changes as the universe expands or contracts.
It is worth noting that the metric given in Eq. (1.1) is left invariant under the following rescalings
a(t)! sa(t), r ! r/s, k ! s2k, (1.2)
where s is a constant. This property, rather conveniently, allows us to rescale the radial coordinate in such a way that the scale factor at present time is equal to unity, i.e. a0 =1.
Observationally it is well-known that a(t)is in fact an increasing function
by noticing that the spectra of distant galaxies are redshifted, i.e. a spectral line with a restframe wavelength lrest is observed to have lobserved >lrest.
This is expected in an expanding universe, as the electromagnetic waves are stretched alongside with cosmic evolution. An important relation between the redshift factor z and the cosmic scale factor a is given by
z⌘ lobserved
lrest 1=
a0
a(t?) 1, (1.3)
where a0 is the present-time scale factor and a(t?) is the value of the scale
factor when the wave has been emitted.
The redshift of galaxy spectra can be interpreted as a result of Doppler effect. When the considered galaxy moves much slower than the speed of light, then the corresponding Doppler redshift of spectral lines would be given by z ⇡ v/c, where c is the speed of light in vacuum and v is the speed of the galaxy with respect to the observer.
As discussed above, in an expanding FLRW universe the physical dis-tances between two points at fixed comoving distance r is given by rphys =
a(t)r. This then leads to the recession speed of a galaxy at the distance rphys
from the observer to be
v= Hrphys ⇡ H0rphys, (1.4)
where H ⌘ ˙a/a, with a dot denoting a derivative with respect to cosmic time t, is known as the Hubble function. In the last part of this equation we have assumed the galaxy to have a small redshift, so that the Hubble function can be assumed to be approximately constant and equal to its present-day value of H0. This result is the celebrated Hubble’s law of cosmic
1.1 the cosmological standard model in a nutshell 7 Dynamics of the FLRW universe
In the context of Einstein’s theory of General Relativity [7], the dynamics of the metric tensor field can be derived from the Einstein-Hilbert action, given by S = M 2 Pl 2 Z d4xp gR+Sm(gµn, Yi), (1.5)
where g is the determinant of the metric tensor, R ⌘ gµnR
µn is the Ricci
scalar constructed from the metric tensor gµn and the corresponding Ricci
tensor Rµn. Sm is the action describing the dynamics of matter fields,
collec-tively denoted by Yi. Additionally, we have introduced the reduced Planck
mass, defined by MPl ⌘
q
¯hc
8pGN, with ¯h being the reduced Planck constant and GN - the Newton’s constant. It should be noted that the central property
of GR is that all the matter species Yi are universally coupled to the metric.
This coupling is proportional to the Newtonian gravitational constant GN.
An additional observation at this point is that the symmetries of the ac-tion (1.5), namely, the invariance under general coordinate transformations, or, the diffeomorphism invariance, allow us to add a constant term in the Einstein-Hilbert action. This term, known as the cosmological constant dis-cussed earlier, is an essential piece for constructing the phenomenologically simplest cosmological model which is compatible with all the currently known experimental and observational evidence, namely the LCDM model. We are going to assume that the matter content of the universe is de-scribed by a perfect fluid with an energy density r(a) and pressure p(a).
Our next step is to derive the equations of motion which govern the dy-namics of this metric. For that purpose we can plug our FLRW metric ansatz Eq. (1.1) into the Einstein-Hilbert action (including the cosmological constant term 2L and the matter energy density r(a)) and obtain the
function N(t) yields the energy constraint equation, which is the celebrated
first Friedmann equation
3M2PlH2 = MPl2 L+r(a) 3MPl2 k
a2. (1.6)
Additionally, the variation of the action with respect to the scale factor a(t) gives ˙H+H2 = 1 6M2 Pl (r(a) +3p(a)) + L 3. (1.7)
where we have set N(t) = 1 and defined the pressure as
p(a) = r(a) 1
3a
dr(a)
da . (1.8)
An important consequence of the diffeomorphism invariance is the auto-matic conservation of the energy-momentum tensor, given the Einstein field equations are satisfied. This conservation is given by rµT
µn =0, where rµ
is the covariant derivative compatible with the metric gµn. For the perfect
fluids considered here this equation takes the form ˙r+3Hr(1+w) = 0,
where w ⌘ p/r is the equation of state of the considered fluid. From this simple relation it follows that the energy densities of dark matter with w = 0, radiation with w = 1/3 and cosmological constant with w = 1
(which are assumed to be non-interacting, hence are conserved separately) are evolving as
rr =rr(a0)a 4, (1.9)
rm =rm(a0)a 3, (1.10)
rL =rL(a0), (1.11)
where a0 is the present-day value of the scale factor.
1.1 the cosmological standard model in a nutshell 9 dominate over radiation at some later stage. Additionally, both radiation and non-relativistic matter will eventually become subdominant compared to cosmological constant. This shows that in the LCDM model the uni-verse asymptotically approaches an epoch described by a constant Hubble function. This spacetime metric at this epoch is known as the de Sitter metric.
It is also useful to introduce the dimensionless density parameters as Wr(a)⌘ rr(a)/3H2M2Pl, (1.12)
Wm(a)⌘ rm(a)/3H2MPl2 , (1.13)
WL(a) ⌘ L/3H2, (1.14)
Wk(a)⌘ k/H2a2. (1.15)
In terms of these dimensionless parameters the first Friedmann equation can be rewritten as
Wm(a) +Wr(a) +WL(a) +Wk(a) = 1. (1.16)
Let us mention that the cosmological observations tightly constrain the spatial curvature k to be tiny [8]. In this thesis we will mainly assume it being exactly zero.
Perturbing the FLRW universe
As we mentioned above, the FLRW metric provides a valid description of the universe on scales larger than O(100) megaparsecs. On smaller
However, it is a fortunate property of the universe that at large enough scales the perturbations of the relevant fields are small enough, so we can make use of the perturbation theory. The starting point for this perturbative approach is to specify the form of the perturbed metric. Naively, one would start perturbing all the components of the metric tensor, which would lead to extremely complicated calculations. However, as we mentioned earlier, one of the central properties of GR is its invariance under general coordinate transformations. For a given calculation in the framework of GR we can choose a particularly suitable coordinate system, where the given problem is solved the easiest. This coordinate freedom is known as the gauge freedom of GR, and the particular coordinate choice is often called a gauge choice for the metric.
In the previous subsection, when deriving the Firedmann equations, we did not make direct use of the Einstein’s field equations. For deriving the equations of motion for the perturbed quantities we can proceed similarly and first derive the action which would then directly lead to the equations of motion for the desired perturbation variables. For example, if we are interested in the linear order perturbations, then we would need to ex-pand the Einstein-Hilbert action to second order in these perturbations. Such a second order action then would lead to linear equations of motion. Alternatively, we could derive the full equations of motion and perturb them to the desired order. In the bulk of this thesis we have used both of these approaches. Here, in order to demonstrate the main features of the standard cosmological model at perturbative level, let us make use of the latter approach.
The starting point are the Einstein’s field equations, derived from Eq. (1.5) by varying with respect to the metric tensor. They read as
1.1 the cosmological standard model in a nutshell 11 where Tµn is the energy-momentum tensor of the matter fields defined as
Tµn ⌘ p2 g dSm
dgµn. (1.18)
The first step of our perturbative treatment is to write the metric as
gµn = ¯gµn+dgµn, (1.19)
where ¯gµn is the FLRW background metric and dgµn is a perturbation
around it. We will then plug it in the left hand side of Eq. (1.17) and keep only the terms up to first order in dgµn. Such a background-perturbation
splitting is an arbitrary choice, but is perhaps the most intuitive one from the point of view of a generic observer in a Hubble flow. The most general form of the metric is
ds2 = (1+2f)dt2+2aBidtdxi+a2 dij hij dxidxj, (1.20)
where one can show that f, Bi and hij are, respectively, 3 scalar, vector and tensor. It turns out that the perturbative calculations simplify significantly if we decompose these perturbations into scalar, vector and tensor degrees of freedom. For the vectors this decompositions is well known from general physics. Namely, any 3 vector can be written as
Bi =∂iB+Si, (1.21)
where B transforms as a 3 scalar, while Si is a divergence-free 3 vector. Similar decomposition is possible for higher-rank objects, namely for hij:
hij =2ydij+2E,ij+Fi,j +Fj,i+˜hij, (1.22)
where y and E are two additional 3 scalars, Fi is a devergence-free
3 vector and the 3 tensor ˜hij is such that ˜hi
As such, we have decomposed the 10 independent components of the symmetric 4⇥4 metric dgµn into 4 scalar functions (namely, f, y, B, E), 4
vector modes (encoded in the 6 components of Bi and Fi and the
corre-sponding divergence-free conditions), and 2 tensor degrees of freedom (encoded in the 6 components of ˜hij and the corresponding conditions given in Eq. (1.23)).
The significant advantage of such a decomposition is that it turns out that the linearized Einstein’s equations lead to decoupled dynamics of these scalar, vector and tensor sectors. The formation of the large scale structure of the universe is largely given by the scalar sector of the metric, and now we will be considering only this sector. Let us mention, however, that the dynamics of the tensor sector characterizes the propagation of gravitational waves, and hence, even though not relevant for the large scale structure formation, contains valuable information by its own.
The most general way to write the scalar-perturbed metric is as follows ds2 = (1+2f)dt2+2a∂iBdtdxi+a2⇥(1 2y)dij+2∂i∂jE⇤dxidxj. (1.24)
A widely used gauge choice is the Newtonian gauge, specified by E =
0= B.
Our next step is to include the perturbed energy-momentum tensor Tµn.
The latter for a generic perfect fluid can be written as
Tµn = (r+P)uµun+pgµn, (1.25)
where, uµ is the four-velocity of the fluid element as seen by a comoving
observer, r is its energy density, p - its pressure. Here we will assume any deviations from the perfect fluid approximation to be exactly zero. The perturbed sector of the energy-momentum tensor is given by
dT00 = dr, (1.26)
dT0i = dTi0 = (1+w)¯rvi, (1.27)
1.1 the cosmological standard model in a nutshell 13 Here we have denoted the spatially averaged energy density as ¯r, and the
perturbations around this background are denoted by dr ⌘ r(x) ¯r(t).
Additionally, vi are the components of the three-velocity and c2
s ⌘ dp/dr
denotes the square of the sound speed of the considered fluid.
The linearly perturbed Einstein equations have the following form (see e.g. Ref. [9]) 6H2f 2 a2∂i∂iy+6H ˙y = 1 M2 Pl dT00, (1.29) 2∂i( ˙y+Hf) = M12 Pl dT0i, (1.30) ¨y+3H ˙y+H ˙f+ (3H2+2 ˙H)f+ 1 3a2∂i∂i(f y) = 1 6M2 Pl dTii, (1.31) 1 a2∂i∂j(y f) = 1 M2 Pl dTij, i 6= j. (1.32)
These equations are more conveniently studied in the spatial Fourier space, i.e. using the spatial Fourier components of the corresponding vari-ables. Our convention for Fourier decomposition for a field j(x)is
j(x) =
Z
d3kjkeik·r, (1.33)
where k is the spatial Fourier wavenumber and r is the spatial real-space coordinate.
Now, going to Fourier space and combining Eqs. (1.29) and (1.30) we obtain the Poisson equation
k2 a2y= 1 2M2 Pl (3H(1+w)¯rv dr), (1.34)
Additionally, for matter sources which have dTi
j = 0 we have an important
relation
f = y. (1.35)
For simplicity in our analysis we will consider only the modes which are very deep inside the Hubble horizon, i.e. k2/a2 H2. Additionally, we will be considering the so-called quasistatic regime, where one assumes that the cosmological variables can change only at the time scales close to the order of the Hubble rate, i.e H2dj ⇠ H ˙dj ⇠dj¨ . In this approximation we have
2k2 a2 y= 1 M2 Pl dr (1.36)
Besides the Einstein equations an extra information is contained in the perturbed conservation equations. The n = 0 and n= i components of the
continuity equation rµTµn = 0 in sub-horizon limit, during dark matter
domination, yield
d0+q = 0, (1.37)
q0+Hq k2 f+c2sd = 0, (1.38)
where we have now started to use the conformal time, related to the cosmic time through adt = dt, and primes denote derivatives with respect to
conformal time. Additionally, we have defined q ⌘ ∂ivi and d ⌘ (r(x)
¯r(t))/ ¯r(t). From these two equations we then obtain the master equation
for linear structure formation
d00 +Hd0+ ✓ c2sk2 3 2H2 ◆ d =0. (1.39)
The perturbations will experience a growing force by gravity, but the growth will be slowed down by the non-zero sound speed (i.e. by pressure). For cold dark matter the sound speed is negligible, c2
sk2 ⌧ H2, and the
1.2 observations 15 1.2 observations
In the past decades several types of observations have become sufficiently robust and now serve as the basis for our current understanding of the cosmological standard model. Let us here briefly discuss the main of these cosmological observables (see, e.g., [10]). Before doing that, however, it is important to mention that various distance definitions are used for interpreting different cosmological observations. Distances between two points in FLRW spacetime are in fact not uniquely defined, so let us start by defining various useful distances and give the relationships among them.
• Comoving distance. The comoving distance from us to an object at a given redshift z is given by
Dcom ⌘ c a0H0 Z z 0 d˜z E(˜z), (1.40)
where H0 is the value of the Hubble rate at present time, and E(z) ⌘
H(z)/H0.
• Luminosity distance. For a source with an absolute luminosity L,
observed to have a flux F on our detectors, we can define the so-called luminosity distance to the source
D2lum ⌘ L
4pF. (1.41)
• Angular diameter distance. For an object of proper size (in the direc-tion perpendicular to the line of observadirec-tion) D`, observed to subtend
an angle Dq, we can define the so called angular diameter distance to the object as
For a spatially flat universe the luminosity distance is related to the comoving cosmic distance by Dlum = (1+z)Dcom. This expression is rather
generic and holds for almost any cosmology. It should, however, be kept in mind that it will be violated in a theory where the photon number is not conserved, for example, due to mixing of photons with some hidden sector. Additionally, the luminosity distance is related to the angular diameter distance by Dlum = (1+z)2Dang.
After this prelude we can start discussing the main cosmological obser-vations.
Supernovae Type Ia. Perhaps the best-known cosmological constraints are from Supernovae. The luminosities (or the absolute magnitudes) of these objects are known to be highly correlated with the widths of their light-curves. This fact allows for an accurate determination of the absolute magnitude, given the light-curve observation of a supernova. A key relation for cosmological purposes is the relation between the distance modulus µ (the difference between the apparent and absolute magnitudes) and the luminosity distance
µ =5 log Dlum/10pc, (1.43)
Having the distance modulus measurements of supernovae, one then can measure the luminosity distance, and hence constrain a particular cosmo-logical model.
Cosmic Microwave Background. The Cosmic Microwave Background (CMB) is one of the major sources of information in cosmology. As we mentioned earlier, after inflation the universe was filled with a hot photon-baryon plasma. The photon-baryons tend to cluster through gravitational attraction, but the photonic pressure stops this clustering. As a result the cosmic plasma experiences acoustic oscillations. When the universe cools down sufficiently the photons decouple from baryons and start to free-stream through the universe. This decoupling happens at redshift zdec ⇡ 1090 and
1.2 observations 17 make up the CMB sky. The fluctuations of the photon temperature are
sensitive to the density perturbations of the relevant energy components at the decoupling era, their velocities and the gravitational potentials. These fluctuations, measured as a function of direction ˆn, can be decomposed in spherical harmonics as
dT(ˆn)
T =
Â
`Â
m a`mY`m(ˆn), (1.44)where T is the average temperature of the CMB, a`m’s are the corresponding angular modes and Y`m(ˆn)’s denote the spherical harmonics.
For each mode`, the variance of the a`m modes is known as the angular
power spectrum, given by C` = 2 1
` +1
Â
m h|a`m| 2i (1.45)
A relatively simple information in the CMB angular power spectrum is encoded in the scale of acoustic oscillations. The effective sound speed of the photon-baryon fluid cs determines this scale through
rs(zdec) =
Z • zdec
d˜zcs(˜z)
H(˜z). (1.46)
The associated angular scale qang(zdec)and the angular diameter distance
to the acoustic scale Dang(zdec)are related with each other by (1+zdec)Dang(zdec) = rs(zdec)
qang(zdec). (1.47)
Baryon Acoustic Oscillations (BAO). The acoustic oscillations men-tioned in the context of CMB affect not only photons but also baryons. Sim-ilarly to the acoustic scale in Eq. (1.46), there is a similar scale for baryons, imprinted during the so-called drag epoch, taking place at zdrag ⇡ 1020,
One expects an enhanced galaxy population at the scales of cosmic struc-ture separated by rs(zdrag). The corresponding angular scale at a particular
redshift z then serves as a useful probe for the cosmic background. The relevant geometric expression is similar to (1.47) and is given by
(1+z)qs(z) = rs(zdrag)
Dang(z). (1.48)
Growth of structure An instrumental quantity often employed in LSS studies is the growth rate f , defined as
f ⌘ dlnddlna. (1.49)
There is a useful fitting formula for this quantity, given by f = Wgm, where
the power is constant in LCDM and is approximately equal to g ⇠ 0.55. An observed deviation from this value will be a smoking gun evidence for beyond LCDM physics.
Weak lensing. One of the striking predictions of any modern theory of gravity is the light deflection by massive sources. Cosmologists have come up with a beautiful idea which exploits the gravitational lensing for measuring the properties of the large scale structure. When a light from a galaxy is travelling trough the LSS, it gets slightly distorted. The distortions of this light can be characterized by the gradient ∂qi
source/∂qj, where qj is
the angle under which we observe the given light ray, while qi
source is the
unaltered (unlensed) angle, and the indices(i, j) label two directions on the
sky. In a theory of gravity (not necessarily GR) this gradient is given by
∂qisource ∂qj dij ⌘ Z rsource 0 d˜r ✓ 1 ˜r rsource ◆ ˜r(f y),ij, (1.50)
where rsource denotes the comoving distance to the considered galaxy.
This matrix is conventionally written as
1.3 the inflationary paradigm 19 It can be shown that the so-called convergence kwl describes the overall
magnification of the sources, while the components of the shear g1 and g2
describe its distortions. The measurements of these quantities and their cross-correlations provide valuable cosmological information.
1.3 the inflationary paradigm
In the previous sections of this introduction we have presented the main ideas of the cosmological standard model. In that discussion we have taken the observed large-scale homogeneity and isotropy of the universe, as well as the small value of the spatial curvature, as granted. They are, however, rather unnatural in the standard FLRW universe with a sequence of radiation and matter dominated epochs. This has motivated the birth of the inflationary paradigm.
Let us start our discussion from the so-called flatness problem. In a deceler-ating universe1 the absolute value of the curvature contribution in Eq. (1.16)
increases, because its denominator aH = ˙a decreases, unless the curvature
of the universe is exactly zero. The observed spatial flatness then suggests that in the past the universe has experienced a phase of accelerated expan-sion, known as cosmic inflation (see e.g. [10] for a pedagogical introduction to inflation).
Another striking issue with the standard cosmological picture is the overall homogeneity of CMB. It can be estimated that the CMB patches of more than ⇠1 degree apart never would have time to communicate with each other starting from the time of infinitely small universe (the Big Bang)
to the time of recombination [10]2. While the CMB photons were, in fact, in
causal contact after the last scattering, the entire idea of CMB suggests that they shouldn’t interact, hence they cannot thermalize after decoupling.
The crucial quantity for our discussion here is the comoving particle horizon, defined as dH,com ⌘ Z a 0 d˜a ˜a 1 ˜aH(˜a), (1.52)
which measures the maximum distance the light could have travelled in FLRW spacetime between times characterized by scale factors 0 and a. It is instructive to rewrite dH,com in terms of the comoving Hubble radius(aH) 1
as dH,com = Z lna 1 dln˜a ˜aH(˜a). (1.53)
The last expression suggests a solution to the horizon problem. If we could have en epoch during which (aH) 1 is increasing towards the past,
then dH,com could be made larger. What we are seeking for is a mechanism
which would make dH,com much larger than (aH) 1 during the standard
expansion. This is precisely the idea of inflation; make the comoving Hubble radius larger in the past, so that the entire observable CMB would have been in causal contact at some point in the past.
It is easy to notice that achieving such a regime does not only resolve the issue with the horizon, but also resolves the flatness problem. Indeed, the condition d(aH) 1/dt <0 implies that ¨a >0 and hence, using the second
Friedmann equation, that w < 1/3, which is exactly the condition for the
flat universe to be an attractor under cosmic evolution.
The most common dynamical realization of inflation is through a canoni-cally normalized scalar field j with a potential V(j). The homogeneous
1.3 the inflationary paradigm 21 and isotropic equation of motion (i.e. taking j(x) to be a function of time
only) of such a field is given by
¨j+3H ˙j+V(j),j =0, (1.54)
where V(j),j ⌘∂V(j)/∂j.
Additionally, the energy and momentum of the scalar field can be shown to be
rj = 12 ˙j2+V(j), (1.55)
pj = 12 ˙j2 V(j), (1.56)
respectively.
In order this field to be able to successfully drive the inflationary dy-namics, we need the equation of state of the scalar field to be close enough to 1, which is the case of the shallow potentials. Additionally, in order to have long enough inflation, we need the above condition to be satis-fied for long enough period. These two conditions are formally written as e ⌘ ˙H/H2 ⌧ 1 and h ⌘ ˙e/He ⌧ 1. The inflationary stage should be followed by the stage of hot FLRW expansion, which means that the inflationary stage must end eventually. This additionally means that infla-tion cannot be realized via a cosmological constant, because in that case the universe would have no physical clock specifying when the inflation should end.
invariant, with a slight tilt characterized by a slope of ns 1, where ns
is typically referred to as scalar spectral index. Additionally, inflation also predicts presence of primordial gravitational waves, again from the initial quantum fluctuations of the metric. The amount of produced primordial gravitational waves are typically characterized by the ratio of powers in tensor and scalar fluctuations, referred to as the tensor to scalar ratio, and denoted by r.
The idea of inflation is summarized in Fig. 1.1. The decreasing comoving Hubble radius (Ha) 1 makes the particle horizon at the epoch of CMB
formation larger compared to the value in the standard, non-inflationary scenario. This additionally resolves the flatness problem. Moreover, and perhaps more importantly, inflation is an elegant mechanism for generating the observed large scale structure from the primordial quantum fluctuations of the inflaton field and the spacetime metric.
1.4 beyond the standard model: dynamical dark energy and modified gravity
1.4 beyond the standard model 23
a
= 1
(aH)
1 Radiation Domination InflationObservable LSS and CMB scales
Cosmic past
Comoving
scales Horizon exit
Horizon entry
Figure 1.1:The idea of inflation is to modify the expansion history of the universe in such a way that the comoving Hubble radius is decreasing before the standard expansion regime starts.
learn about the fundamental theories. In practice this argument, of course, is more complicated, as the observational consequences of, for example, string theory so far are rather ambiguous. An interesting development in this direction was suggested in [13]. The authors have conjectured that the scalar field potential for all consistent theories should satisfy the constraint
|rfV|
This conjecture is in contrast to the string theory landscape scenario [14–19] (see Ref. [20] for a brief review of related ideas), where it is considered that string theory describes an enormous number of metastable de-Sitter vacua.
There is no consensus about the theoretical validity of Eq. (1.57) in the string theory community (see [21] for a review). Moreover, the use of the conjecture in its current shape for cosmological phenomenology is still rather ambiguous. Indeed, as the conjecture does not specify the value of the constant c, it is difficult to confront it with phenomenological studies. The way forward in this situation is to study the phenomenological implications of the models presented in [13] which have been served as the primary support for the conjecture. These models are given in terms of concrete potentials and therefore their precise phenomenologies can be worked out. The main result of such an investigation in [22] is that all these considered models are incompatible with cosmological data. This, perhaps, is difficult to interpret as a very strong observational challenge for Eq. (1.57) because, again, in the latter the imprecise nature of c makes it impossible to draw decisive, quantitative conclusions.
Even if not making the connection to any fundamental theory, the dy-namical dark energy scenario is still very interesting to study. One of the interesting motivations to study such alternatives is the so-called coincidence problem, which is based on the question of why is the universe starting to ac-celerate exactly at the present time, and not, say, much later in the future. In this context let us discuss a particularly appealing feature, namely the pres-ence of the so called scaling fixed points in the phase space of quintesspres-ence models. The model that we will study has a simple exponential potential of the form V(f) = V0elf. For l2 >3(wB+1), with wB being the equation of
state for the background fluid (e.g. dark matter or radiation), the universe enters a scaling regime where the scalar field mimics the evolution of the background fluid, with wDE =wB; the dark energy density parameter takes
the form WDE =3(wB+1)/l2 (see [23] for a review). This scaling property
1.4 beyond the standard model 25 (chosen to bep750 in this example), where we have shown the evolution of
the quintessence energy density compared to that of dark matter. The figure shows that the scalar field, after some oscillations, quickly follows the back-ground and one can achieve a scaling solution during matter domination in this example. Te horizontal axis here is N ⌘ lna, with N =0 corresponding
to the present time. Such scaling solutions may indeed provide a solution to the coincidence problem. Even though these are very interesting features, the obvious problem, of course, is that a single-exponential potential has a constant slope, and therefore, once the scaling regime is switched on it never ends, so there is no dark energy domination.
A particular extension of the considered model is the following two-exponential potential
V(f) = V1el1f+V2el2f. (1.58)
The phenomenological merit of this double exponential model is that under certain conditions the scaling solution can gracefully exit to the desired accelerating phase at late times. This transition can be obtained if
l21 > 3(wB+1)and l22 <3(wB+1) in the potential (1.58). At early times,
the potential is dominated by the el1f term, for which the scalar field follows the equation of state of radiation and/or matter, hence scaling solutions. Later in the evolution of the universe, the el2f term dominates, for which the evolution is not of the scaling form and the late-time attractor is the scalar field dominated solution (with WDE =1). In this scenario, the asymptotic
value of the dark energy equation of state is wDE = 1+l22/3, providing
viable cosmologies, just as for the single exponential with l2 < 3(w
B+1).
The right panel of Fig.1.2shows an example of this so-called scaling freezing scenario with the double-exponential potential, where the transition from the scaling evolution to the scalar field dominated evolution has been depicted.
-15 -10 -5 0 5 10 10-9 10-7 10-5 10-3 N ρDE /ρ M -15 -10 -5 0 5 10 10-8 10-5 0.01 10 104 N ρDE /ρ M
Figure 1.2:The ratio of the dark energy density rDE to that of matter rM as a function
of N. The left panel demonstrates the scaling solutions of a single-exponential model V(f) = V0elfwith l2 >3(wB+1), while the right panel is for V(f) =
V1el1f+V2el2f with l21 >3(wB+1)and l22<3(wB+1).
There is, however, another exciting prospect. As we mentioned earlier, gravity is the most relevant interaction at the cosmological scales. This means that cosmology is an ideal playground for testing the underlying theory of gravity. In order to effectively study the limitations of GR at cosmological scales, one needs to consider its viable modifications.
GR, in fact, is the unique theory of interacting, massless, spin-2 field (see [24] for a proof). This immediately suggests that in order to construct an alternative to GR one can either consider a massive extension of the latter, or add extra dynamical degrees of freedom, such as additional scalar field(s).
1.4 beyond the standard model 27 A particular problem for a given theory is the presence of unstable solu-tions. Commonly discussed types of instabilities are the so-called ghost and gradient instabilities.
Let us start by discussing the gradient instability. It basically originates from a wrong sign gradient term in the Lagrangian of the theory. For the simplest possible example let us consider a scalar field theory in Minkowski spacetime which has a wrong sign spatial gradient term. The equation of motion for the scalar field j of such a theory in Fourier space is simply given by
¨jk k2j =0, (1.59)
where k is the absolute value of the spatial Fourier wavenumber. Note that in a healthy theory the second term would have been with an opposite sign. The solutions of this equation scale as jk(t) ⇠e±kt, the growing part
of which leads to a gradient instability. The characteristic timescale of the instability scales with the wavenumber as 1/k.
Another widely encountered type of pathology is the ghost instability. To understand ghosts it is enough to consider the following, non-gravitational toy example for two scalar fields c and j
L = 12∂µc∂µc 12∂µj∂µj+V(c, j), (1.60)
where the potential V(c, j) is given by
V(c, j) = 1
2m2cc2
1
2m2jj2+lc2j2, (1.61)
with mc and mj being the masses of the fields and l a positive constant.
infinite [25, 26] (assuming the considered theory is valid up to arbitrarily high energies). This means that the presence of ghosts makes the theory highly undesirable. Ghost fields are typically present in theories whose equations of motion contain higher than second order time derivatives [27,
28]. This fact is one of the main locomotives for constructing alternatives to GR. One of the most well studied class of theories is in fact the Horndeski theory - the theory of a single scalar field coupled to gravity in such a way that the resulting equations of motion are second order in time [29, 30]. This last requirement ensures the absence of ghosts.
Horndeski theory is a generalization of scalar-tensor theories known since a long time ago. One of the first examples is the Brans-Dicke theory [31], the main idea of which is to promote the gravitational constant to a dynamical field. In the so called Jordan frame (which means that the matter fields are minimally coupled to the metric gµn), the Brans-Dicke theory has the
following action S = M 2 Pl 2 Z d4xp g 1 2jR wBD 2j rµjrµj V(j) +Sm(gµn, Yi), (1.62) where wBD is a constant. The GR limit of this theory is recovered in the
limit of infinitely large Brans-Dicke parameter wBD.
1.4 beyond the standard model 29 with ¯j being the homogeneous background sector of the scalar field j,
and m being the mass of the scalar field. As we see, contrary to GR, here the gravitational strength, which controls the effectiveness of dark matter clustering, is a function of scale and time.
Additionally, the relation between the two gravitational potentials (which in GR is given by the simple identity Eq. (1.35)) in the quasistatic limit is given by:
h(a, k) ⌘ f
y =
2(1+wBD) +m2 ¯ja2/k2
2(2+wBD) +m2 ¯ja2/k2. (1.65)
Interestingly, the functional forms of these two new fucntions µ(a, k) and
h(a, k)are generic for the entire class of Horndeski gravity [32]. Particularly,
these can be written as
µ(a, k) = h1(a)1+k 2h 5(a) 1+k2h3(a) (1.66) h(a, k) = h2(a)1+k 2h 4(a) 1+k2h5(a), (1.67)
where hi are functions of background only, and their form is model-specific. Horndeski gravity is expected to be constrained by several high-precision large-scale structure surveys. However, the recent detection of the gravi-tational waves originating from a pair of merging neutron stars and the simultaneous detection of their electromagnetic counterpart, the LIGO event GW170817 [33] and its counterpart GRB 170817A [34], have already cut a large portion out from the Horndeski Lagrangian. This has been achieved through the strong bounds imposed on the speed of gravitational waves (which is constrained to be very close to the speed of light in vacuum); see [35] for a recent review on the topic. Note, however, that the mentioned bound on the speed of gravitational waves is strictly valid only at the scales of LIGO events, which is k ⇠ O(10 100) Hz. Horndeski gravity, on the
present-time cosmic expansion rate, H0, which is about 20 orders of
magni-tude smaller than the LIGO scale. This means that for interpreting the LIGO bounds one might need to include corrections to the considered theories, which can then naturally bring the speed of gravitational waves in these theories to be very close to the speed of light.
Let us conclude this section by mentioning that while the Horndeski-type general approach to Modified Gravity is very fruitful, it still misses some important classes of theories. Among these modifications to gravity, the bimetric theory of ghost-free, massive gravity is of particular interest. It stands out especially because of the strong theoretical restrictions on the possibilities for constructing a healthy theory of this type. Indeed, historically it has proven to be difficult to invent a healthy theory of massive, spin-two field beyond the linear regime. The linearised theory has been known for a long time [36], while at the fully nonlinear level the theory has been discovered only recently by constructing the ghost-free theory of massive gravity [37–46]. This development has also naturally led to the healthy theory of interacting, spin-2 fields, i.e. the theory of ghost-free, massive bigravity [47]; see Refs. [48–50] for reviews.
Over the past decade, there has been a substantial effort directed towards understanding the cosmological behaviour of bimetric models, both theo-retically and observationally. Particularly, it has been shown that bigravity admits FLRW cosmologies which perfectly agree with cosmological ob-servations at the background level (see Ref. [51, 52] for reviews). At the level of linear perturbations the cosmological solutions have been shown to suffer from either ghost or gradient instabilities, although the latter can be pushed back to arbitrarily early times by imposing a hierarchy between the parameters of the theory [53]. It is also conjectured [54] that the gradient instability might be cured at the nonlinear level due to the presence of the Vainshtein screening mechanism (see later in this chapter) in the theory.
1.5 screening mechanisms in modified gravity 31 couple to only one of the two metrics (spin-2 fields). The metric directly
coupled to matter is called physical metric, and the other spin-2 field, called reference metric, affects the matter sector only indirectly and through its interaction with the physical metric.
1.5 screening mechanisms in modified gravity
One of the most well-understood properties of modified gravity theories is that there is an extra (often referred to as a "fifth") force in addition to the standard Newtonian force. To understand this effect let us study the following, quite generic coupling of the matter fields to the scalar field sector: S = Z d4xp g " M2 Pl 2 R 1 2rµjrµj V(j) # +Sm(˜gµn, Yi), (1.68) where ˜gµn ⌘ A(j)gµn. (1.69)
A central equation here is the geodesic equation for a non-relativistic test particle in the Newtonian limit:
¨xi +Gi00 = dlnA
dj rij, (1.70)
where G denotes the Christoffel symbol, and xi are the spatial coordinates
of the considered test particle. This equation motivates us to interpret the right hand side as a fifth force.
for the fifth force to be screened in an environment-dependent manner. For demonstrating the main idea behind the common screening mechanisms let us consider the field equation of motion of the theory given in Eq. (1.68)
⇤j =V(j),j dlnA
dj Tr[Tµn], (1.71)
where Tµn is the Einstein frame metric, ⇤ ⌘ rµr
µ is the d’Alambert
operator, and the trace is taken with gµn. For a non-relativistic matter sector,
such as cold dark matter, Tr[Tµn] = r, with r being the matter density.
This motivates us to define Veff(j; r) ⌘ V(j) +rlnA(j); an effective
potential which reacts to the matter density of the ambient space. Let us discuss two qualitatively different choices for the A(j) function and the
potential V(j):
• V(j) = Ljnn+4, A(j) = ej/Mc,
• V(j) = 12µ2j2+ l4j4, A(j) = 1+2Mj22s,
where L (not to be confused with the cosmological constant), µ, l, Mc and
Ms are constants.
The first of these choices is in the class of Chameleon screening mecha-nisms [56], the idea of which is to enhance the effective mass of the scalar field, hence rendering the corresponding fifth force to be a short-range one; see the left panel of Fig. 1.3. The second choice corresponds to the Symmetron mechanism, [57] the idea of which is to suppress the coupling of the scalar field to the matter; see the right panel of Fig. 1.3.
Another important mechanism is the Vainshtein mechanism [58, 59], which relies on the non-linearities of the scalar field induced due to higher order derivative self-couplings, such us L ⇠∂µj∂µj⇤j. Vainshtein
1.6 the era of precision cosmology 33 V( ) 1 2 Veff( ; 1) Veff( , 2) 0 Veff( ; > M2 2) Veff( ; < M2 2)
Figure 1.3:Left panel: Demonstration of the Chameleon screening mechanism Right panel: Demonstration of the Symmetron screening mechanism. See the text for details.
1.6 the era of precision cosmology
Before summarizing the content of this thesis let us present a comment on how fast the presion of cosmological observations grows. The quality of modern cosmological datasets posits very high standards in front of cosmological model building initiatives. As a striking demonstration of this let us examine Fig.1.4, which shows the current observational constraints on inflationary models by the CMB data given by the Planck collaboration [60] alongside with the same constraints from a decade-old WMAP collaboration [61]. We see that many interesting models, e.g. the polynomial inflationary models with potential fk with k = 2, 4/3, 1, 2/3, are now disfavored or
ruled out by date. All these models were inside the 95% sweet spot of the data in 2009 provided by the WMAP collaboration, as one can see in Fig. 1.4, while they are now either outside or close to the boundary of the 95% confidence region of the Planck 2018 data.
0.93 0.94 0.95 0.96 0.97 0.98 0.99 1.00 1.01
Primordial scalar tilt (ns)
0.00 0.05 0.10 0.15 0.20 0.25 0.30 Tens or- to-sc al ar rati o (r0 .002 ) Convex Concave WMAP, 95% C.L. PLANCK-2018, 95% C.L. 68% C.L. -attractors Hilltop quartic model Natural inflation Power-law inflation V 2 V 4/3 V V 2/3 R2inflation
Low scale SB SUSY
N = 50 N = 60
Figure 1.4:Evolution of precision in inflationary parameters over a decade, from WMAP [61] to Planck [60]. The reconstructed Planck constraints correspond to the combination TT,TE,EE+lensing+BK14+BAO provided in [60]. One can look, for example, at the area between ns = 0.95 and ns = 0.98. Although both of
these values were inside the 68% contour back in 2009, they are now strongly disfavored with more than 95% confidence.
normal in the past, especially in string theory phenomenology. The same concerns such expressions as "parametrically small", or "parametrically large". We can see examples in Fig. 1.4showing that reducing the bound on r from⇠ 0.08 to⇠0.04 has made various theoretical ideas either supported or ruled out by the precision data in cosmology.
Similarly, the constraints on dark energy become more and more pre-cise. For example, 15 years ago the constraints on the parameter l in the exponential potential e±lf for quintessence, allowed l= 1.6 [62,63].
1.7 this thesis 35 difference between numbers such as l < 1 and l < 1.4. Indeed, models
with l > 1 are ruled out by cosmological observations with more than
99.7% confidence, whereas the condition l <1 is not satisfied by the string
theory models of [13]. 1.7 this thesis
• Chapter2is dedicated to a study of a new class of inflationary models known as cosmological a-attractors. We promote these models towards a unified framework describing both inflation and dark energy. We construct and study several phenomenologically rich models which are compatible with current observations. In the simplest models, with vanishing cosmological constant L, one has the tensor to scalar ratio r= 12aN2, with N being the number of e-folds till the end of inflation, and the asymptotic equation of state of dark energy w = 1+ 9a2 .
For example, for a theoretically interesting model given by a =7/3
one finds r⇠ 10 2 and the asymptotic equation of state is w ⇠ 0.9.
Future observations, including large-scale structure surveys as well as Cosmic Microwave Background B-mode polarization experiments will test these, as well as more general models presented here. We also discuss the gravitational reheating in models of quintessential inflation and argue that its investigation may be interesting from the point of view of inflationary cosmology. Such models require a much greater number of e-folds, and therefore predict a spectral index ns
that can exceed the value in more conventional models of inflationary
a-attractors by about 0.006. This suggests a way to distinguish the
conventional inflationary models from the models of quintessential inflation, even if the latter predict w= 1. This chapter is based on
• The topic of Chapter3 is the theory of massive bigravity, where one has two dynamical tensor degrees of freedom. We consider an inter-esting extension where both of the metrics are coupled to the matter sector, which is known as the doubly-coupled bigravity. The main aim of this chapter is the study of gravitational-wave propagation in this theory. We demonstrate that the bounds on the speed of gravitational waves imposed by the recent detection of gravitational waves emitted by a pair of merging neutron stars and their electromagnetic coun-terpart, events GW170817 and GRB170817A, strongly limit the viable solution space of the doubly-coupled models. We have shown that these bounds either force the two metrics to be proportional at the background level or the models to become singly-coupled (i.e. only one of the metrics to be coupled to the matter sector). The mentioned proportional background solutions are particularly interesting. In-deed, it is shown that they provide stable cosmological solutions with phenomenologies equivalent to that of LCDM at the background level and at the level of linear perturbations. The nonlinearities, on the other hand, are expected to show deviations from LCDM. This chapter is based on Ref. [65].
1.7 this thesis 37 the effective radiation and curvature terms be within observational
bounds. The late-time acceleration must be accounted for by a sepa-rate positive cosmological constant or other dark energy sector. We impose further constraints at the level of perturbations by demanding linear stability. We comment on the possibility of distinguishing this theory from LCDM with current and future large-scale structure surveys. This chapter is based on Ref. [66].
• The final Chapter5is dedicated to the study of the effects of screening mechanisms in modified gravity on the dynamics of the spherical collapse of dark matter. In particular, we investigate the splashback scale in Symmetron modified gravity. The splashback radius rsp has been
identified in cosmological N-body simulations as an important scale associated with gravitational collapse and the phase-space distribution of recently accreted material. We employ a semi-analytical approach, namely the self-similar spherical collapse framework, to study the spherical collapse of dark matter halos in Symmetron gravity. We provide, for the first time, insights into how the phenomenology of splashback is affected by modified gravity. The Symmetron is a scalar-tensor theory which exhibits a screening mechanism whereby higher-density regions are screened from the effects of a fifth force. In this model, we find that, as over-densities grow over cosmic time, the inner region becomes heavily screened. In particular, we identify a sector of the parameter space for which material currently sitting at the splashback radius rsp, during its collapse has followed the
formation of this screened region. As a result, we find that for this part of the parameter space the splashback radius is maximally affected by the Symmetron force and we predict changes in rsp up to around
Part I
2
D A R K E N E R G Y, a - AT T R A C T O R S , A N D L A R G E - S C A L E S T R U C T U R E S U RV E Y SThis chapter is dedicated to a study of a new class of inflationary models known as cosmological a-attractors. We promote these models towards a uni-fied framework describing both inflation and dark energy. We construct and study several phenomenologically rich models which are compatible with current observations. In the simplest models, with vanishing cosmological constant L, one has the tensor to scalar ratio r = 12aN2, with N being the num-ber of e-folds till the end of inflation, and the asymptotic equation of state of dark energy w = 1+ 9a2 . For example, for a theoretically interesting model
given by a =7/3 one finds r ⇠ 10 2 and the asymptotic equation of state
is w ⇠ 0.9. Future observations, including large-scale structure surveys as well as Cosmic Microwave Background B-mode polarization experiments will test these, as well as more general models presented here. We also discuss the gravitational reheating in models of quintessential inflation and argue that its investigation may be interesting from the point of view of inflationary cosmology. Such models require a much greater number of e-folds, and therefore predict a spectral index ns that can exceed the value
in more conventional models of inflationary a-attractors by about 0.006. This suggests a way to distinguish the conventional inflationary models from the models of quintessential inflation, even if the latter predict w = 1.
This chapter is based on: Y. Akrami, R. Kallosh, A. Linde, V. Vardanyan,
2.1 introduction 43 2.1 introduction
In this chapter we are going to construct viable dynamical dark energy models in the context of recent progress achived in cosmological applica-tions of supergravity. We particularly will be using some novel ideas which have been discovered in inflationary cosmology. More concretely, recent investigations have found a broad class of theories, known as cosmological
a-attractors, which are based on models where the kinetic term of a scalar
field has a pole [76–81]. In such theories, the potential has a plateau shape, exponentially rapidly approaching a constant at large values of the inflaton field j. These models, to be described in section 2.2 of this chapter, are favored by the recent inflation-related cosmological observations [82].
Because of the extreme flatness of the potential in a-attractors, these models can be suitable not only for describing inflation but also to describe dark energy, see e.g. Refs. [83–88]. Moreover, it may also be possible to find
a-attractor models which can simultaneously describe inflation and dark
energy [84,87, 88] in the context of the quintessential inflation [89].
In this chapter, we extend the investigation of the quintessential inflation models based on a-attractors. We study models with arbitrary L, relax some of the assumptions made in Refs. [84,87, 88], and consider a much more general class of theories. In particular, we describe the a-attractor version of the simplest linear dark energy model, a model with exponential potential with two shoulders proposed in Ref. [90], and a generalized version of the model studied in Refs. [84,88].
The asymptotic value w• of the parameter w in the equation of state
p = wr for quintessential inflation depends on the limiting value of the
quintessence potential. If this value is negative, the universe eventually collapses, but under certain conditions it may pass through a temporary but long stage of acceleration. Here we call w• the asymptotic value of
equation of state wDE and the observable "all-inclusive" effective equation
of state weff.
If the potential V of the quintessential inflation models asymptotically vanishes (i.e. if the cosmological constant is zero), the value of w• in the
simplest models is given by w• = 1+ 2
9a . (2.1)
Interestingly, the difference between w• and the equation of state w = 1
for the cosmological constant is inversely proportional to a, whereas the tensor to scalar ratio is directly proportional to it,
r = 12a
N2 , (2.2)
where N corresponds to the remaining number of e-folds from the end of inflation at the moment of generation of perturbations studied by WMAP and Planck. This may help us either to rule out, or to confirm theories of that type by a combination of searches for B-modes and investigation of dark energy.
Note that this result is valid only if the cosmological constant is zero, which provides us with an intriguing possibility to test this hypothesis. Meanwhile in the theories with a negative cosmological constant, the uni-verse eventually collapses. However, in some cases one may have a pro-longed state of accelerated expansion, just as in the model proposed in Ref. [91].
If the asymptotic value of the potential is positive (i.e. if the cosmological constant is positive), and the quintessence field slowly rolls towards infinity, the universe asymptotically approaches a de Sitter regime with
2.1 introduction 45 This is the most general regime that is relatively easy to achieve in the su-pergravity constructions discussed here. Of course, if these models correctly describe our world, the observations looking for deviations of quintessence from the cosmological constant will not bring us anything exciting. But there may be a silver lining here.
Indeed, the process of reheating in the models of quintessential inflation is non-standard, and it can be very inefficient. In that case, the inflaton field after the end of inflation may enter a long stage when its energy density is dominated by the kinetic energy with w = +1. This simple fact affects
the number of e-folds N [84]. Indeed, as we will show, the number of e-folds in the a-attractor models of quintessential inflation with gravitational reheating can be greater than the corresponding number in the conventional (non-quintessential) versions of a-attractors and in the Starobinsky model by DN ⇠ 10. This is a significant difference, which may have important observational consequences.
In particular, the general prediction of a attractors for ns is
ns = 1 N2 . (2.4)
One can easily check that the difference between ns for conventional
a-attractors with N ⇠ 50 and a-attractor models of quintessential inflation with N ⇠60 is about 0.006, which coincides with 1s error bar in the Planck 2015 results [82]. This increase in the value of ns and N is not very easy to
achieve otherwise, see e.g. Refs. [92, 93].
This suggests that future observations may be able to differentiate be-tween the regular versions of inflationary a-attractors and their quintessen-tial generalizations. More generally, we might be able to differentiate, though somewhat indirectly, the cosmological constant and quintessence without relying on extreme accuracy in measuring w. This is a rather intriguing byproduct of the present investigation.
